Download ESS200C Lecture 12 Solar Wind Interactions with Unmagnetized

ESS154/200C
Lecture 10
Solar Wind Interactions:
Unmagnetized Planets
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ESS 200C Space Plasma Physics
ESS 154 Solar Terrestrial Physics
M/W/F
10:00 – 11:15 AM
Geology 4677
Instructors:
C.T. Russell (Tel. x-53188; Office: Slichter 6869)
R.J. Strangeway (Tel. x-66247; Office: Slichter 6869)
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Topic
Instructor
A Brief History of Solar Terrestrial Physics
CTR
Upper Atmosphere / Ionosphere
CTR
The Sun: Core to Chromosphere
CTR
The Corona, Solar Cycle, Solar Activity
Coronal Mass Ejections, and Flares
CTR
The Solar Wind and Heliosphere, Part 1
CTR
The Solar Wind and Heliosphere, Part 2
CTR
Physics of Plasmas
RJS
MHD including Waves
RJS
Solar Wind Interactions: Magnetized Planets YM
Solar Wind Interactions: Unmagnetized Planets YM
Collisionless Shocks
CTR
Mid-Term
Solar Wind Magnetosphere Coupling I
CTR
Solar Wind Magnetosphere Coupling II;
The Inner Magnetosphere I
CTR
The Inner Magnetosphere II
CTR
Planetary Magnetospheres
CTR
The Auroral Ionosphere
RJS
Waves in Plasmas 1
RJS
Waves in Plasmas 2
RJS
Review
CTR/RJS
Final
Due
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PS2
PS3
PS4
PS5
PS6
PS7
Interactions with Atmosphereless Bodies
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If a body in a flowing magnetized space
plasma is nonconducting or only partially
conducting, like rock or ice, some or all of
the incident flow is absorbed on the
plasma ram face.
The absorbed plasma leaves a wake behind
the moon. If the magnetic field is parallel
to the flow, the wake does not close (fill in)
completely farther downstream. If the
magnetic field is perpendicular to the floe,
the wake closes along the field at the
thermal speed. The field in the wake is only
modestly distorted from the undisturbed
external field in both cases.
The external flows may be solar wind or
magnetospheric, and the bodies involved
may be planetary satellites or asteroids.
Examples of such interactions include the
Earth’s moon, Tethys, Rhea, Dione at
Saturn. Note that sputtering of the surface
by the incident particles may also form a
rudimentary atmosphere, usually an
exosphere, on the plasma ram face, but we
ignore this here.
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Interactions with Atmosphereless Bodies
• When the body has a low conductivity surface layer but an interior ‘core’ of
both significant size and conductivity, the electrical currents induced in that
conductor can persist sufficiently long to exclude the magnetic field from the
conducting region.
• If the conducting region is large compared to the size of the body then a very
large disturbance can be produced (above left), and sometimes detected.
• This effect can be used to sound the size of the core and has been used on our
Moon. It can be measured when the Moon enters the lobes of the Earth
magnetotail where external fields are quiet and only slowly varying.
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Interactions with Bodies with Weak
Atmospheres
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Bodies with weak atmospheres or exospheres
may lose mass to the solar wind flow via photoionization or impact ionization or charge
exchange followed by the ion pick-up via the
Lorentz force. As described in the solar wind
lecture, the pickup ion drifts at the plasma flow
speed in a cycloid perpendicular to the magnetic
field if the plasma flow and field are
perpendicular.
The momentum required for the pick-up
particles is extracted from the ambient flow,
whether solar wind or magnetospheric. This is
often referred to as ‘mass loading’ of the plasma
because the newly produced ions usually add
mass to the original flow (the exception is for a
few special cases such as H+ charge exchange
with H).
Since gyro-radii are mass-dependent, a mass
spectrometer effect occurs in which some ions
with small gyroradii compared to the body size
may impact it while others manage to avoid
impact. The latter can escape if their speeds
exceed the escape velocity for the body,
determined by its gravity and size. The impact of
these atmospheric ions can cause additional
sputtering loss.
Bottom: illustration of pickup
Ion impact on Earth’s Moon
from Manka and Michel, 1973
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Interactions with Bodies with Thick
Atmospheres and No Magnetic field
• When a body has a substantial atmosphere, and is sufficiently close to the Sun,
solar EUV photons partially ionize the dayside to produce an ionosphere.
• When the external plasma flow arrives, it could impact the atmosphere and be
absorbed if there is no magnetic field. But if the incident plasma is magnetized,
currents are induced in the conducting ionosphere that oppose the external
magnetic field and counter its penetration (diamagnetic currents).
• Eventually (months?) these currents would decay, but if the direction of the field
keeps changing, the field is excluded. What is happening is δB/δt induction.
These ionospheric currents form a magnetic barrier that can exclude the
external magnetic field from the ionosphere and lead to a magnetosheath-like
structure with draped interplanetary fields. This type of interaction is often
referred to as an ‘induced magnetosphere’. Venus is our best example of this. 6
The Ionospheric Obstacle: Ionopause
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The ionosphere is the obstacle to the solar wind. The boundary between the
ionosphere and the shocked and slowed solar wind is called the ionopause.
Around solar maximum, the ionopause is located at about 400 km altitude at the
subsolar point and around 1000 km at the terminator.
The ionosphere thickness is quite small compared to the planet radius, and so the
obstacle is almost planet-sized and much smaller than the magnetosphere of Earth
(Venus’ twin in planet size).
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Size of the Ionospheric Obstacle
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The solar wind flow exerts dynamic pressure plus magnetic and thermal pressure. The ionosphere exerts
mainly thermal pressure force against the solar wind. Along the stagnation line of the interaction,
approximate pressure balance is achieved.
As the incident solar wind changes, or the ionosphere thermal pressure changes (e.g. due to solar EUV
variations) the ionosphere moves. It goes to lower altitudes as the external pressure increases but has a limit
at around the exobase altitude (at Venus ~200km). Below this altitude the ionosphere is photochemically
controlled, meaning production rather than dynamics dictates the thermal pressure.
If the thermal pressure at the peak of the ionosphere is greater than the dynamic pressure of the incident
flow, and the standoff distance is well above the collisional region of the atmosphere (e.g. above the exobase),
then the magnetic field is largely excluded from the ionosphere. If the peak ionospheric thermal pressure is
lower than the external plasma pressure, the magnetic field collisonally diffuses into the ionosphere to
contribute to its pressure and reinforce its obstacle-like effects on the oncoming flow. Such ‘magnetized
ionospheres’ were first seen at Venus.
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The Bow Shock Location
• The Venus Bow Shock position is observed to change in response to solar EUV
flux (Left figures). It is also expected that the ionosphere is generally more
magnetized but we lack observations of the solar minimum subsolar
ionosphere.
• The general idea of how the Venus interaction changes from solar maximum
to minimum in the form of the draped magnetic field geometry, is suggested
in the right-hand sketch above. However Venus Express observations around
solar minimum indicate the actual geometry may involve both configurations,
depending on the hemisphere relative to the convection electric field
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direction.
The Magnetic Barrier
External solar wind pressure
nswkTsw + ρv2 + B2 / 2μ0
Is mainly dynamic
Dynamic pressure
goes to
Magnetic pressure
goes to
Ionospheric thermal
Pressure (magnetic
contributions can
remain)
Along the stagnation line pressure balance occurs and a transition is observed
from solar wind dynamic pressure to magnetic pressure in the inner draped
field/magnetosheath region, to the ionosphere where thermal plasma pressure
dominates. The layer where the draped magnetic field pileup takes most of the
pressure has been referred to as a ‘magnetic barrier’ or pileup region (from Russell
and Vaisberg in VENUS, Univ of Arizona Press, 1982)
Flux Ropes in Planetary Ionospheres
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The large scale fields are absent when the ionopause is high (~ 300 km subsolar), or
decay if they were present. In some cases, typical for solar maximum, bundles of
magnetic flux resembling flux ropes are observed in the dayside ionosphere. These
fields (e.g. from weak downward diffusion or decayed fields) become force-free so
that the twist in the field balances the magnetic pressure gradient.
In a force-free rope, the current is parallel to the magnetic field. If j = αB and α is
constant, this is called a Taylor state, and the field components (axial and azimuthal)
are described by Bessel functions.
The bottom panel above right illustrates the distribution of flux ropes in the Venus
ionosphere, viewed in cross section. In this sketch the barrier and axial fields point
out of the page. However some analyses suggest they are often twisted and kinked. 11
The Magnetosheath
The magnetosheath fields at Venus fit models remarkably well (here a gas-dynamic
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field model was used). The ionosphere acts like a conducting sphere in this case.
Numerical Simulations of Venus 1
MHD simulations can replicate many
average properties of the Venussolar wind interaction. The solar
wind hydrogen plasma and the
ionospheric plasma are the main
‘fluids’ represented.
The simulations use an atmosphere
Model (Venus has a mostly CO2 and
O atmosphere) with ions produced
by ‘exposing’ the atmosphere to
solar EUV (photo-ionization), impact
ionization by solar wind electrons,
and charge exchange with solar
wind protons.
Sample of PVO
magnetic field
strength averages
folded into a plane.
Dark blue areas are
not sampled. Here
the Sun is at left.
The model results show the
magnetic field strength (contour) and
directions (white arrows) (left panels)
and flow speed and directions (right
panels) (right panels). As in the real
case, an induced magnetospheric
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obstacle is produced.
This figure shows both the upstream and downstream views of the interaction process. The
gray isosurface represents a density contour with planetary O+ density equals to 50/cc as an
indication of Venus ionosphere, which is around 400 km altitude along the subsolar line, but
extends nearly to 5 RV in the nightside. This figure shows how field lines wrap around the
obstacle on the upstream side, but slip over the obstacle on the night side.
The field lines at high latitude are kinked because the plasma closer to the planet have been
mass-loaded by planetary ion production and so significantly slowed down as compared with
plasma further away from the planet.
The draped field exerts a JxB force on the plasma in the wake, speeding it up to escape
velocity as down-tail distance increases.
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The MHD simulations of the Venus
solar wind interaction can be used
to study the ion escape associated
with thermal pressure gradient ,
convection electric field –VxB, and
JxB forces. Here the loss of O+ in
the wake is illustrated by color
contours of the ion density in the
equatorial plane views for solar
max and min cases. (The
interplanetary field orientation is
equatorial)
The escape velocity for Venus is
~11 km/s so all ions moving
tailward above this speed are
permanently lost to the planet.
The rates are ~3-5X1025 s-1.
Such escape processes acting over
time can affect the planet’s
atmosphere evolution.
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Atmospheric Erosion 1- Superthermal Atoms
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The Venus exosphere has a hot hydrogen and a hot oxygen component. The hot
oxygen is produced by the dissociative recombination of the primary ion O2+
e + O2+ → O* + O*
(here * indicates a neutral O atom with supra-thermal energy from the process)
Some of these suprathermal atoms are upward-directed and reach high altitudes.
Like any thermal atoms above the ionopause, these atoms can be ionized by
photoionization, impact ionization and charge exchange and can add to the
production of ions directly exposed to solar wind convection electric fields. These
ions are then subject to the ion pickup process.
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Atmospheric Erosion 2-Loss due to ion pickup
ion trajectory
if Vsw=0
Convection E
field=-Vsw x B
accelerates
the ion here
B field
Vsw
Incident Plasma
flow
Pickup Ion
trajectory
Because the solar wind
speed is high and even the
shocked solar wind in the
magnetosheath has fast flows
(see MHD model velocities)
the ions that are picked up
can reach velocities above the
~11 km/s escape velocity.
Recall how the ion pickup process
works. One can think of it as ions
becoming entrained on moving field
lines. (Note in the lefthand picture,
in a frame moving at Vsw the ion is
simply gyrating around the field.) In
this case the moving magnetic field
is provided by the solar wind, and
ions by the atmosphere and various
ionization processes.
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Atmospheric Erosion-Observations and Models
Cross-flow IMF
PVO Plasma Analyzer
detections of inferred
escaping O+ ions organized
by the Interplanetary
Magnetic Field (IMF)
direction show the spatial
asymmetry expected for
the ion pickup process,
characterized by a
preference for the –VxB
electric field hemisphere.
Test particle pickup ions
showing the expected
asymmetries for IMF in
east and west orientations.
These are based on
magnetic fields and flows
from an MHD model.
An ionospheric obstacle with crustal fields
The Mars-solar wind interaction represents an ionospheric interaction but
with further complications compared to Venus:
1)There are significant crustal remnant magnetic fields present.
2)The gyroradius of the incident solar wind protons is more comparable to
the scale of the interaction features (e.g. the subsolar magnetosheath),
making microscopic or kinetic effects more important overall.
A few field lines traced in the
Purucker model of the Mars
crustal magnetic field (only).
the surface is color coded by
the radial field strength (Red
and blue are strong +/- field).
The strongest fields are in the
southern hemisphere.
Crustal field intrusions in MGS Magnetic field data and their apparent effects on the
magnetosheath field draping show the need for more complicated models for Mars
The Ionosphere of Mars
• The Mars and Venus
ionospheres, illustrated by
the photochemical models
shown in the upper figure
panels are similar but not
identical.
• Although both are mainly
O2+ at the ionosphere peak,
Venus has higher densities of
atomic oxygen ions above
the peak.
• At both planets, the
observed ion density at high
altitudes is less than
expected, consisted with
erosion/escape at the top
(bottom panels).
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Mars’ Atmospheric Loss
Mars is smaller than Venus and has lower surface gravity. While dissociative
recombination-enabled escape of O (O2+ + e ->2O*) doesn’t work on Venus
(need 10 eV (vesc=11 km/s O)), it dominates all other O escape processes for
Mars (needs only 2 eV (vesc=5 km/s O)).
Ionospheric ‘source’
From Shinagawa, JASR, 2005
Numerical Simulations for Mars
Sophisticated numerical models are necessary to understand the complicated geometry of
the Mars solar wind interaction and to interpret the phenomena observed in the plasma
interaction. The figures show calculated magnetic field and velocity in the meridianal plane
for a typical case. The color plots show the magnitudes; the white lines marked with arrows
indicate the vector direction of the magnetic field and the arrows show the direction (not the
magnitude) of the velocity. The dashed line represents the mean bow shock and the dashdot line is the mean MPB locations from Vignes et al. [2000]
Ma et al. model ionosphere at various locations vs. Viking data
SZA=45
SZA=0
SZA=45
This self-consistent MHD model with an ionosphere reproduces the
Viking Lander profiles as well as the bow shock position.
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Sub-Alfvenic, steady field interactions:
Alfven Wings
When the magnetic field is strong
so that the flow is sub-Alfvenic, the
field lines bend but do not strongly
drape.
If the flow across the polar cap
becomes very slow perhaps due to
intense mass loading, then the flux
tube is essentially frozen to the
moon (e.g. Io) and the other flux
tubes have to move around the
Alfven wing field lines rooted to the
moon. The extensions of the
penetrated flux tubes are called
‘Alfven wings’. They can be
detected far from the body itself, at
increasing distances from the wake
axis.
Jupiter’s satellite Io was the
inspiration for work on this type of
interaction although many
magnetosphere-satellite
interactions at the giant planets are
likely of this type.
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Alfven Wing plus Ionosphere Interactions
Saturn’s large (~2275 km radius) satellite Titan has a thick atmosphere that
makes it like a small planet. It typically resides in Saturn’s magnetosphere although
at ~20 Rsaturn orbital distance it sometimes crosses the magnetopause into
Saturn’s magnetosheath. The Cassini mission is exploring its plasma interaction.26
Titan’s orbit (red) is near the
Magnetopause and goes in and
out of the bowl-like plasma disk
Complications of the Titan interaction include a changing geometry between
the sunlit face of Titan where the main ionosphere is produced, and the oncoming
magnetospheric flow that corotates with Saturn. Also the local field is not so dipolar,
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and proximity to magnetopause and magnetotail make its environment variable.
Titan’s ionosphere
Titan’s upper atmosphere is mainly molecular nitrogen and methane- while
the ionosphere is a rich mix of hydrocarbon ions. Thus it does not have the
same photochemistry as Venus and Mars.
(Cassini INMS profiles of atmospheric neutrals
from Waite et al., Science, 2005)
(modeled dayside ionosphere
From Ma et al., JGR, 2007)
Single-Fluid MHD Model Results in the Ideal Case
U
U
U
B0
E0
B
U-E Plane
U-B Plane
The top two panels show plasma flow speed(contour color) and direction
(white arrows) in both U-E and U-B planes. The bottom two panels show
magnetic field strength (contour color) and direction(white arrows) in both UE and U-B planes. The single fluid MHD model results are symmetric in UE
plane because in this model, the mass-loading effect is included but the
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gyroradii of the pick-up ions are neglected.
Multi-Fluid MHD Model Results in the Ideal Case
U
U
U
B0
E0
B
U-E Plane
U-B Plane
The multi-fluid MHD model allows the motions of different ion fluids
to be decoupled, and it was able to reproduce the asymmetric
results along the convection electric field direction (similar to hybrid
model results). Multi-fluid MHD model also predicts that more
heavy ions are escaping from Titan.
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SLT Effect– 18 SLT vs. 6 SLT
u0
E0
Sun
Sun
The figures show the density distribution(contour) of two different
mass fluids: mass 1(upper panels) and mass 28(lower panels) over
plotted with the corresponding flow vector(the white arrows) in the
UE plane. The plasma interaction also depends on Titan’s orbital
location as shown in the two cases: 18 SLT(left panels) vs. 6
SLT(right panels).
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Interactions including Plumes and Volcanoes
Enceladus’ Plume
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The interiors of some moons are outgassing rapidly with much of that gas escaping.
– Io has volcanoes.
– Enceladus has a southern polar plume from fissures that vent mainly water ice and dust.
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Enceladus’ plume extends far south of it so the central mass-loading point is well below
the moon.
MHD simulations with charge-exchange, impact and photoionization can mimic the
variation in the magnitude of the perturbation but not the components, indicating that
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there are forces and flows we are not properly calculating.
The Cometary Interaction
• When comets approach the Sun,
they heat up and outgas.
• The expanding gas decreases in
density and is lost through
photoionization
Q
Qo
R
exp(
)
2
R
u
where u is the outflow velocity
(~km/s) and τ the photoionization
time scale.
• The comet can produce a small
field-free region around the nucleus
where cometary ions dominate , but
the region of draped interplanetary
magnetic field due to cometary ion
pickup producing solar wind mass
loading is huge.
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Cometary Simulations
• A numerical simulation of this process shows that the stream lines of the
flow do not become very diverted but flow almost straight through the
mass-loading region. Note the scales.
• The field lines become draped and eventually straighten far downstream.
• Comets also produce lots of dust. This dust is charged and can interact
with the solar wind like very heavy pickup ions.
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Space Weather Effects at Unmagnetized Bodies
Interplanetary
magnetic field
associated with
an ICME
Periods of
Venus obstacle
Encounters
Once/day
Many ICMEs and SEP events are observed going past Venus (here an ICME
seen on PVO) and Mars, but the of study their effects (e.g. on atmosphere
escape) has yet to be carried out in depth. This is a project for the coming years.
Summary and Conclusions
• The interaction of flowing magnetized plasmas with
atmosphereless bodies can tell us about the surface and
interior of the body.
• If there is a weak atmosphere, it can be probed using the IMF
to produce a mass spectrum.
• A thick atmosphere can support an ionosphere which can be
sufficiently electrically conducting that the interplanetary
magnetic field is excluded from the ionosphere. This forms a
magnetic barrier that deflects the solar wind and forms a
shock. A magnetotail forms. There may be large scale fields or
flux ropes in the ionosphere.
• Some moons outgas rapidly in plumes and volcanoes and also
form mass-loading obstacles. Charge-exchange producing fast
neutrals is also important.
• Comets are active outgassing bodies that mass load the solar
wind. Although the smallest bodies, comets can produce some
of the largest unmagnetized plasma interactions in the solar
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system.